Abstract

Quasar outflows may play important role in the evolution of its host galaxy
and central black hole, and are most often studied in absorption lines. In
this paper, we present a detailed study of multiple outflows in the obscured
ultra-luminous infrared quasar Q1321+058. The outflows reveal themselves in
the complex optical and UV emission line spectrum, with a broad component
blueshifted by 1650km s−1and a narrow component by 360km s−1, respectively.
The higher velocity component shows ever strong N iii]
(N iii]/C iii]=3.8± 0.3 and N iii]/C iv=0.53)
and strong Si iii] (Si iii] /C iii]≃ 1), in addition
to strong [O iii]λ5007 and [Ne iii]λ3869 emission.
A comparison of these line ratios with photoionization models suggests an
overabundance of N and Si relative to C. The abundance pattern is consistent
a fast chemical enriching process associated with a recent starburst,
trigerred by a recent galaxy merger. The outflow extends to several tens
to hundred parsecs from the quasar, and covers only a very small sky. We
find that the outflow with line emitting gas is energetically
insufficient to remove the ISM of the host galaxy, but total kinetic energy
may be much larger than suggested by the emission lines. The velocity range
and the column density suggest that the outflow might be part of the low
ionization broad absorption line region as seen in a small class of quasars.

The optical and UV continuum is starlight-dominated and can be modeled with
a young-aged (1 Myr) plus an intermediate-aged (∼0.5−1 Gyr) stellar
populations, suggesting a fast building of the stellar mass in the host
galaxy, consistent with the starburst-type metal abundances inferred from
the high velocity outflow spectrum. The broad band spectral energy
distribution shows that it is an obscured quasar with its bulk emission
in the middle infrared. The star formation rate, independently estimated
from UV, far-infrared, and emission line luminosity, is much lower than
that is required for the co-evolution of the black hole and its host
spheroid.

Vigorous starburst and nuclear activity will result in a strong
negative feedback to the host galaxy. Massive winds have been observed in
starburst galaxies, ULIRG and AGN with mass outflow rates from 10 to
1000 M⊙ yr−1(e.g. (Rupke et al. 2005 ); (Lípari et al. 2005 )). It was proposed
that such a feedback will quench gas supply to the AGN and the
star-formation region, and result in a sudden halt of the
star-formation process in the galaxies (e.g. (Di Matteo et al. 2005 )). The
passively evolved galaxies will first appear as E+A in the
subsequent ∼ 1 Gyr, and then as red galaxies. The latter can
explain the red colors of massive early type galaxies ((Springel et al. 2005 )).
If quasar activity is going-on for sometime after the
star-formation essentially ceased, the object will show
characteristics of Q+A. Q+A’s have been detected now with a large
number in the Sloan Digital Sky Survey ((York et al. 2000 ), SDSS).
Most of these systems show weak [O ii] emission, likely from
the Narrow Line Region (NLR), indicating that major star-forming
activity has ceased((Zhou et al. 2005 )), or, at least, is greatly
suppressed. Indeed, Ho (2005) found that the ratio of SFR to mass
accretion rate in PG quasars is well below that required for
sustaining the MBH−σ
relation for nearby spheroid galaxies. It was further suggested
that this is due to that star formation is suppressed rather than
that cold gas is exhausted. It should be noted that the author used the
[O ii] emission-line luminosity to infer SFR, for which the reddening
correction is largely uncertain, however.

Gas outflows are thought to play an important role in the galactic scale
feedback. Outflows on scales of tens to hundreds parsecs have been observed
in nearby Seyfert galaxies and radio galaxies with velocities of a few
hundreds to thousand km s−1(e.g. (Gabel et al. 2005 ); (Das et al. 2005 ); (Ruiz et al. 2005 )). However,
modeling the absorption and emission lines for well-studied nearby Seyfert
galaxies suggested a relative low mass loss rate and small kinetic power
in these outflows (Crenshaw, Kraemer & George 2003). Even higher velocity outflows,
from a few 103 to 104 km s−1, are detected in broad absorption line
(BAL) quasars (e.g. (Weymann et al. 1991 )), but the scale and the total mass in
BAL outflows are still not clear, neither the relation between the
BAL region (BALR) and the narrow/broad emission-line region (N/BELR).

The recent intensive starburst in the circum-nuclear region has also
significant implications for the gas phase metallicity. Massive stars
evolve very fast and produce a chemical pollution in the circum-nuclear ISM
primarily with α-elements via SN ii, while less massive stars evolve
much slowly and make their chemical contribution of iron peaked elements,
mainly through SNIa, with a delay of at least 1 Gyr. The detail abundance
pattern depends on both the history of star-formation and the stellar
initial mass function. Supersolar metallicity has been suggested
by analyzing the broad line and broad absorption line spectra
(e.g., Hamann 1997; Baldwin et al. 2003)
, which are produced in a relative
small region enclosed the active nucleus.
However, it is still not clear whether this reflects an over-all high
metallicity or only these elements concerned.

In this paper, we study in detail the optical and ultraviolet (UV)
spectra and the broad band properties of the ULIRG Q 1321+058.
This object was classified as a quasar based on a low resolution
spectrum taken in the optical identification processes of HEAO 1-A2
(The High Energy Astrophysics Observatory-1) X-ray sources ((Remillard et al. 1993 )).
However, it was not detected by XMM-Newton ((Bianchi et al. 2005)), which
led these authors to propose that it is a Compton-thick, type II object.
The optical image taken with Hubble Space Telescope (HST) showed a strong
disturbed morphology possibly with a double nucleus of separation less
than 1 kpc, suggesting a recent merger of two gas-rich galaxies
((Boyce et al. 1996)). Darling & Giovanelli (2002) detected an OH Megmaser
with a multi-peaked profile, which they interpreted as possible multiple
masser nuclei.
Here we show that its broad band properties are consistent with an
obscured quasar. The optical and the UV continuum is likely dominated
by starlight, that is also typical for a type II object.
Extreme velocity outflows (EVOFs) were noticed previously in the
[O iii] and Hβ emission-lines (Lípari et al. 2003 ). Taking the
advantage of the broad wavelength coverage of the SDSS spectrograph,
we have identified the EVOFs in Hα and [Ne iii]λ3869. We
find that the EVOFs also appear in the UV lines in the HST (Hubble
Space Telescope) FOS spectrum, such as C ivλ1549, N iii]λ1750,
Al iii]λ1860, Si iii]λ1892, and C iii]λ1909.
The line ratios imply a dense line emitting gas with a fast metal
enriched history. We find with interest that the velocity range and the
column density of the AGN-driven outflows are comparable to those in
low-ionization BALR
(LoBALR), implying that the object would appear as a LoBAL quasar if
our line-of-sight passing through the outflows.
These unusual properties of Q 1321+058 provides us with an ideal
laboratory to study the impact of AGN and starburst feed-backs to
galaxy formation and evolution.
Throughout this paper, we
assume a Λ-CDM cosmology with
H0=72km s−1 Mpc−1, ΩΛ=0.7 and
Ωm=0.3.

The SDSS spectrum used in this work was extracted from the
SDSS archive2. The UV spectrum was
observed with the high-resolution gratings of the Faint Object
Spectrograph (FOS) on board HST (Bechtold et al. 2002).
We retrieved the calibrated 1-D spectrum from the HST archive
using the on-flight calibration. The spectra of two exposures
are combined to remove cosmic rays.
Both optical and UV spectra are wavelength- and
flux-calibrated.
The spectra are corrected for
Galactic reddening of E(B-V)=0.031 ((Schlegel et al. 1998 )) and
are shifted into the rest frame of the object using a
redshift of z=0.20467 as
measured from optical emission lines (see below).

2.1 Continuum

The optical continuum shows prominent high order Balmer absorption
lines, thus is likely dominated by starlight (Fig 1). We
model the continuum with the templates derived using the ICA method
(refer to Lu et al. (2006) for detail) assuming no contribution
from the AGN:

f(λ)=A(λ)∑aiICi(λ)

(1)

where A(λ) is the dust extinction factor, ICi the ith
independent component. Note that both ICi and ai are positive.
The intrinsic reddening is taken as a free parameter and the extinction
curve of Calzetti et al. (2000) for starburst galaxies is adopted. The
templates are broadened and shifted
, up to 1000 km s−1 in velocity space, to match the observed
spectrum. With the redshift noted above, we find that
Δv=−40±30 km s−1. Emission line regions were masked during the
fit. The best-fit starlight continuum is shown in Fig 1. The
presence of Balmer absorption lines and lack of an apparent 4000Å
suggest a recent/on-going starburst. An effective reddening of
starlight is yielded to be E(B−V)=0.65±0.05 for the Calzetti
et al. ’s extinction curve, which is among the most dusty galaxies
in the SDSS galaxy sample((Lu et al. 2006 )).

Considering the relative low S/N ratio of the HST-FOS spectrum, we fit
the UV continuum with either a reddened simple stellar population (SSP)
or a reddened power-law. Prominent emission lines are modeled with
Gaussians in this fit with their widths and centroids as free parameters;
these include C iv]λ1549, Si iv+O iv]λ1400, He iiλ1640,
O iii]λ1663, N iii]λ1750, Al iiiλ1860, N ivλ1486
Si iii]λ1892, and C iii]λ1909. Weak emission lines
detected only marginally, such as Si iiλλ1808,1817,
are not considered. The prominent emission line centered at
λ1782.4 rest frame, which can be identified as blue shifted
Fe ii 191 (wavelength 1787Å), is modeled with a Gaussian. In the worst
case, wavelength zero point of FOS can be displaced up to
250 km s−1 (Keyes, HST Handbook), we will keep this in mind in the
following analysis of the FOS data.

The SSP spectra fSB99(λ) was taken from starburst99 models
((Leitherer et al. 1999 ); (Vázquez & Leitherer 2005 )). The high resolution SSP models are only available
for the solar abundance with a wavelength coverage of 1200−1870\AA,
and with ages from 1 to 20 Myr for either an instantaneous burst or a
continuous star formation. To extend to the full wavelength range of
the FOS spectrum, we use the interpolated low resolution model of
Starburst99 for λ>1870Å. This is reasonable since the
extrapolated wavelength range is rather small (Δλ=50\AA)
and free from strong stellar absorption lines. The UV continuum can be
reasonably well fitted with this simple model (Fig 2). We use
both SMC and Calzetti’s extinction curves. But the best fitted model
with Calzetti’s extinction curve predicts an
optical flux well above the SDSS spectrum, thus will not be discussed
further. The best fitted model converges to a stellar
population (1 Myr) and a moderate reddening E(B−V)∼0.26±0.03
for the SMC extinction curve. The SFR rate is 450±50 M⊙ yr−1
for instantaneous burst and 430±40 M⊙ yr−1 for continuous
formation models after scaled to a lower stellar mass cutoff of
0.1 M⊙. For instantaneous model 2 Myr population yields significant
worse fit (Δχ2=10), however, for continuous formation model,
2 Myr population also yields reasonable fit (Δχ2=3) with a
factor of two smaller SFR 230±20 M⊙ yr−1. Note that only
statistical uncertainty is quoted here, while much large errors may be
introduced by our assumptions, such as uniform dust reddening and ignoring
intermediate age stellar populations. Thus the SFR value must be interpreted
with care.

The optical spectrum, obtained by subtracting the best UV model, can
be fitted with a reddened stellar population of an intermediate age
around 0.90 Gyr (Fig 2), with a reddening of E(B−V)≈0.30 using Calzetti et al.’s extinction curve. The derived stellar
mass is 9×1010 M⊙. Using the SMC extinction curve, we
obtain a similar reddening (E(B−V)=0.40) and a slightly younger
stellar population (0.64 Gyr) with a somewhat smaller total stellar mass
(13% less). With those exercises, albeit with some uncertainties,
we conclude that a few ∼1010 M⊙stellar mass has been
built within the last Gyr in the host of Q 132+058. Note the older
stellar population has a negligible contribution to the UV flux.

The power-law model yields a fit to UV spectrum as good as the SSP
model. Therefore, we
cannot distinguish two models from the UV fit alone. If both power-law
index and reddening vary freely, they are poorly constrained. For
an unreddened power-law, a slope of β=2.0±0.2 (fλ∝λβ) is obtained. Fixed the slope to typical quasar value
β=−1.7, a reddening of E(B−V)=0.28±0.02 is yielded for
SMC-type grain. However, extrapolating these models over-predicts
severely the optical flux below 4000Å. Only when β≤−6.0,
the model flux is consistent with the optical flux with a large intrinsic
reddening E(B−V)≥0.605. Not only such a flat UV spectrum has never been
observed in quasars, but also the reddened corrected UV flux around 1400Å
is a factor of 3600 times the observed value, which gives an UV luminosity
well larger than the total infrared luminosity estimated
in §4.2. Thus we believe that a single power-law model is
not realistic.

2.2 Emission lines

We obtain the emission-line spectrum by subtracting the model continuum.
The emission-line profiles are rather complex and show several distinct
emission-line components (see Fig 3):

C1

The low-ionization forbidden lines,
[O ii]λλ3726,3729,
[O i]λ6302, [S ii]λλ6717,6732 and
[N ii]λ6584
display a single-peaked profile at a redshift of z=0.20467, which can be
taken as the systematic velocity. There is a corresponding peak in
Hβ, [O iii]λ5007. However, this component is absent in
the ultraviolet lines.

C4

The isolated UV lines C ivλ1549 and
N iii]λ1750 show
a single-peaked profile, which can be well fitted with a Gaussian
blueshifted by ∼ 1,800 km s−1 relative to the systematic redshift.
This component is also evident in [Ne iii]λ3869, and
[O iii]λ4959 and in [O iii]λ5007,
which is blended
with the C3 of [O iii]λ4959.
The extended blue wing of Balmer
lines, and C iii], Si iii] can be attributed to
this component, too.

C2

[Ne iii] shows an additional component centered at the velocity
around -500km s−1. A peak at this velocity is also evident in
[O iii]λ5007, Hβ line. It is likely the weak blue
wing in the [O ii]λ3727, [S ii] and [O i]
is due to this component, too. A comparison of Hα line profile
with those of low ionization lines such as [S ii], [O ii],
[O i] and [N ii] also suggests such a component in
Hα.

C3

Both C iii], Hβ and [O iii] show extended red wing up
to a velocity of ∼1,500km s−1, which is not present in [O ii],
[S ii], C iv and N iii. Note that the redside of
[Ne iii]λ3869 is affected by both He i and
[Fe vii]λ3890, thus it is not
clear if this component is also present in [Ne iii].

C1 and C4 are well defined from the isolated lines such as C iv,
N iii], [S ii], [O i], and [O ii]. One gaussian
can adequately describe these components. C2 and C3 are less cleanly
defined due to their blending with other components. It is interesting
to note that the velocity separation between C1 and C2 is very close
to the separation (490km s−1) of two peaks detected in the OH maser
profile (Darling & Giovanelli 2002). In the following fit, we will
use a gaussian for all components. The good fit to the
C iii]+Si iii] blending with only C3 and
C4 components indicates that C3 can be approximately described with a
gaussian. While a gaussian C2 component yields also reasonable fit to
the optical lines suggests that such approximation is acceptable.

Note in passing, the prominent emission line at 1782.4±0.2Å has
a width of σ=387±34 km s−1 and flux of (189±14)×10−17 erg cm−2 s−1 (see also the lower panel in Figure 2).
The line width is significantly broader than those of C1 and C2, but much
narrower than those of C3 and C4. The line can be identified as
blended of Fe ii multiplet 191 (1785.26, 1786.74 and 1788.07Åwith
intensity ratios 20:20:18). After taking into account of different
instrumental resolution and line blending, the line
center and width are consistent with the C2 component.

In order to obtain more physically meaningful fit, we tied the center
and width of each component to be same for different lines. The flux ratio
of multiplets is either fixed at its theoretical value or set as a free
parameter. The [O iii]λ4959/λ5007 is fixed at its
theoretical value of 1/3 for C1, C2 and C3, and is allowed to vary for
C4. N iii]1750 consists of a blend of five lines at 1746.82,
1748.65, 1749.67, 1752.16 and
1754.00Å, and their relative ratios are fairly constant at density below
107.5 cm−3. We fixed the branch ratios at 0.03, 0.096, 0.502, 0.355,
0.097 as predicted by CLOUDY (See Appendix 1). C iv has two multiplets with a
separation of 500 km s−1. The multiplet ratio (1548/1551Å) is constrained
in the range 2.0 to 1.0.
We assume that λ1900 blending features are dominated by Si iii]
and C iii], rather than by their forbidden counterparts. Justification for
these assumptions can be found in Appendix 1. To couple with the systematic
wavelength calibration uncertainty of FOS, which can be up to on diode
(or 250 km s−1), we allow all UV lines to be systematically
shifted up to 250 km s−1. We assume that each line in the [N ii] and [S ii]
doublets has the same profile. The [N ii]λ6583/λ6548 ratio
is set to its theoretical value of 3/1. The O iii]λ1660/λ1666
doublet ratio is fixed to 1/2.4 for C3 and C4. The initial values for the
centroid and width of each component are estimated by fitting individual
lines with a Gaussian: [S ii], [O i], [O ii], and [N ii]
for C1; [Ne iii], [O iii]
and Hβ for C2 and C3; and C iv, N iii], and [O iii] for C4. A total of
19 emission lines are fitted simultaneously, including C iv,
N iii],
Al III, Si iii], C iii], [O ii], [Ne iii], Hγ+[O iii]λ4363,
Hβ, [O iii]λλ4959, 5007, [O i]λ6302,
Hα+[N ii]λλ6548,6584, [S ii]λλ6717,
6731, He iiλ1640, O iii]λλ1660,1666,
N ivλ1486 and FeII 191.
The final fit is shown in Fig. 3.

Balmer decrements are very different for different components from above
decomposition: C2 has the largest Hα/Hβ ratio of 12, following
by C1 (6.2), C3 (5.1) and C4 (4.2). Both of them are considerably steeper
than 3.1 for an unreddened AGN. If this is interpreted as
due to dust reddening, then C2 is heavily obscured while the others are
mildly extincted. In order to see if this is caused by the decomposition
procedure, we plot Hα/Hβ at different velocity bin. Here,
Hα profile is obtained by subtracting the best-fit [N ii] model
from the observed spectrum. Because C1 is well defined and C2 is weak for
[N ii], [N ii] subtraction should not introduce substantial
uncertainty. The result is shown in Fig 5, where the data are
adaptively rebined to ensure the S/N ratio in each bin. Obviously, the
Balmer decrement varies dramatically across the profile: it rises slowly
from -4000 to -1000 km s−1; then increases sharply from -800 km s−1 and
reaches its maximum of Hα/Hβ∼10 at the C2 centroid
and then decreases rapidly to the red wing. Thus large Balmer decrement
of C2 is independent of the detailed line decomposition. The well defined
profile of Balmer decrement provides an additional evidence for C2 as
a distinct component.

We have decomposed empirically the emission lines into four components.
In this section, we will try to give a physical interpretaion for those
components based on the physical conditions that are required to
reproduce the observed emission line ratios. We will use the photoionization
models as a guideline in the following discussion, but keep in mind that
shocks may also play certain roles giving the presence of outflows indicated
by blue-shifted line profiles.

3.1 Component 1

C1 has doublet ratios of [S ii]λ6716/[S ii]λ6731≃ 1.5
and [O ii]λ3729/[O ii]λ3726 ≃ 0.70, which reach their
low density limits, and hence the density of the emitting gas should be no larger than
a few hundred cm−3. However, its classification is ambiguous on the conventional
BPT diagram ((Baldwin et al. 1981); (Kewley et al. 2006)). On the [N ii]/Hα versus [O iii]/Hβ
diagram (Fig 4), C1 is very close to the extreme H ii curve of Kewley et al. (2001).
It is in the regime for LINER on [O i]λ6303/Hα vs
[O iii]/Hβ diagram, but on the H ii side on the [S II]/Hα
vs [O iii]/Hβ diagram. The Balmer decrement
Hα/Hβ≈6.1±1.6 indicates substantial reddening
(E(B−V)=0.61±0.26 mag, assuming an intrinsic Hα/Hβ=3.0).
With such a reddening, the weakness of the corresponding component in the
UV lines can be explained naturally. With an observed ratio
[O ii]λ3727/[O iii]λ5007≃7.3, the
reddening-corrected ratio would be around 14, extreme among the LINERs
((Kewley et al. 2006)). This puts it on the border between H ii and LINER
on the [O i]/Hα vs [O ii]/[O iii] diagram.
This component can be a mixture contribution from H ii and AGN,
a combination of photoionization and shock-excitation, or gas ionized by
a diluted AGN continuum (Dopita & Sutherland 1996).

3.2 Component 2

At first glance, C2 is formed in the star forming regions. [N ii],
[S ii], and [O i] emission lines is relatively weak with
respect to Hα. Those line ratios fall well in the H ii
locus on all three BPT diagrams (Fig 4). The gas has a relative
low metallicity (Z∼0.2−0.3 Z⊙), according to
[N ii]/Hα ratio ((Pettini & Pagel 2004); Nagao et al. (2006)). However, giving the
high luminosity of the galaxy, the low gas metallicity is not expected.
In addition, large
[Ne iii]λ3869/[O iii]λ5007≃1.15 requires
a rather high gas density (>107 cm−3). Even higher density is required
if the large Balmer decrement of C2 is attributed to dust reddening. In this
case, weakness in [N ii], [S ii] and [O i] is naturally
explained because they are collisionally de-excited, and BPT diagrams no long
provide diagnostics for the ionizing source, as such the ionizing source can
be AGN as well.

The strong Fe ii 191 emission in component 2 is a puzzle. First, its
emission region cannot be reddened as severely as indicated by the apparent
Balmer decrement, which suggests an intrinsic dust reddening of E(B-V)=1.2
mag using Calzetti’s extinction curve for an intrinsic Hα/Hβ=3.0.
This indicates either a large intrinsic Hα/Hβ or a separated
Fe ii 191 emission region. Second, because of lack a similar
component in other UV lines, Fe ii 191 must be produced in a region
shielded from intensive Uv ionizing radiation. While whether this can be
consistent with the prominent emission in [O iii] and [Ne iii]
needs to be further studied, a large column density may explain this. Third, Fe ii
191 is strong while optical Fe ii emission is weak. The dominance of
emission from high excitation levels indicates that either it is formed in a
high temperature and dense HI region or/and through UV photon excitation
by the continuum radiation (PCR). A high temperature and dense HI region
may help to explain Hα/Hβ ratio in this component through
collisional enhancement of Hα emission.
Because this component is blueshifted by 350
km s−1 with respect to the systematic velocity, interaction between
the outflow and interstellar medium may be a viable heating source.
PCR process had been suggested for the Fe ii emission from the wind
of spectroscopic binary KQ Puppis (Redfors & Johansson 2000). In order
to make photon-pumping effective, the gas must see a strong soft UV
photons and has large internal velocity gradient. We speculate that an
ionizing continuum filtered by a thin partial ionized source may mimic
the soft UV photons (see a similar idea in Collins et al. 2009). The
PCR process would also produce Fe ii 193 emission lines at 1459,
1465, 1474Å but the spectrum quality does not allow us to assess
their existence.

Therefore, we tentatively interpret this component as from a dense region
: The large Balmer decrement is
interpreted as due to radiation transfer effect or caused by collisional
excitation in the warm dense HI region, rather than dust reddening;
Fe ii 191 is produced either in a high temperature warm HI region
or through PCR process; [N ii], [S ii] are suppressed by
collisional de-excitation. The blueshift of C2 (350 km s−1) relative to
the systematic velocity can be readily interpreted as from an
outflow with the far-side, redshifted counter-flow being obscured.
However, we cannot rule out the possibility that C2 consists of emission
lines from a dusty star forming region and an additional dense warm partially
ionized gas component.

3.3 Component 3

C3 has strong Si III], C III], Al III, [O iii], [Ne III] and Balmer lines,
but is undetected or very weak in C iv, N iii], [O I],
[O II], [S II] and [N II]. Relatively high ionization and excitation suggest
that this component is produced by AGN photoionization. O iii]1665/[O iii]5007,
C iii]1909/[O iii]5007 and [O iii]5007/Hβ suggest a density of
107.5−8 cm−3, while [O iii]λ4363/λ5007 and
[Ne iii]3896/[O iii]5007 indicate a somewhat lower gas density about
107 cm−3. As [O iii]4363 is badly blending with Hγ line and
the flux of [Ne III] is sensitive to the way of modeling the red wing, the
latter two should be used with care. Therefore, we believe that
108 cm−3 is more likely the real value.

However, the large line ratio of Si iii]/C iii]∼1 can
only be reproduced with a high density nH>109.5 cm−3, if the gas has
abundances being the solar or the scaled solar values (See Appendix 2).
This may be explained in terms of at least two emission line regions
with different densities. Such models, however, have certain drawbacks as the
following. If the high and low density gases are at the same distances
and see the same continuum, the large difference in density would result in
an ionization parameter for the low density gas being three orders of magnitude
higher than that for the high density gas; this would lead to strong coronal
emission lines, such as [Fe VII], which are not seen in the observed spectrum.

A single zone model can be constructed if Si element is enhanced relative to
C. Overabundance of α-elements relative to C is predicted by the
starburst chemical evolution models for quasar host galaxies (Haimann &
Ferland 1993). Considering the fact that Q1321+058 experienced a recent
starburst (see next section), this chemical enrichment route is plausible.
In their models, more massive stars are produced with respect to
star formation in the Galaxy, thus leads to both fast metal enrichment
and a high final metal abundance (≥ 10 Z⊙). α-elements are
enhanced relative to C because C-enrichment from intermediate mass stars
has a finite time delay with respect to
SN II+Ib eruption, which produces most α-elements. It should be
noted that the relative abundances of different metals depend on the
metallicity. With increasing metallicity, N relative to O increases
very fast at all abundances, while Fe and α-elements rise
substantially. Lack of N iii], N iv] and C iv, we can only
put an upper limit on the gas metallicity to 3 solar value based on
N iii]/C iii] because other lines are not sensitive to the
metallicity.

Over the density range considered in Appendix 2, Si iii]/C iii]
provides a good
diagnostics of the ionization parameter. The observed Si iii]/C iii]
ratio can be reproduced with two branches of ionization parameters: an upper
branch with logU∼−0.5 and a lower branch −2.0≤logU<−1.5,
at a column density logNH(cm−2)=21. At larger column
densities, we find that only the low ionization branch is possible, however.
In the density range, C iii]/C iv is sensitive both
to the ionization parameter and to the column
density. The lower limit on C iii]/C iv suggests a large column density
logNH≥22 and a low ionization parameter logU<−2.2, thus the
lower branch solution is preferred. With nH∼107.5 cm−3,
logU∼−2.0 and logNH≥22, the observed line ratios can be
reproduced roughly. The relatively high density can also explain
naturally the weakness of this component in [O ii], [N ii],
and [S ii].

With a relative large emission line width and close to systematic velocity,
this component may be gravitational bound optical thick clouds ionized by
the AGN continuum. We speculate that it is the intermediate emission line
region between BLR and NLR.

3.4 Component 4

C4 is present in all permitted and semi-forbidden lines, as well as
the high-ionization forbidden lines [O iii] and [Ne iii].
The line ratios [O iii]4363/[O iii]5007,
O iii]1665/[O iii]5007,
C iii]1909/[O iii]5007, and
Hβ/[O iii]5007 give consistent densities around 107cm−3,
while [Ne iii]3869/[O iii]5007 suggests a somewhat lower density (106.6
cm−3). These line ratios are only weakly dependent on the metal
abundances or gas column density, thus we believe the density is quite
robust. As for C3, at this density range, a large Si iii]/C iii]
requires that Si is over-abundant relative to C, and this may be
interpreted as the starburst abundances of Haimann & Ferland (1997).
Large N iii]/C iii] ratio can be reproduced only with a very supersolar
abundance (Z∼10Z⊙; see Fig 7) because nitrogen abundance
increases with the average metallicity (Z) as ∝Z2 at high
Z in their models.

As seen in Fig 7, we are unable to reproduce quantitatively all line
ratios with a single constant density models described in Appendix B,
but we believe this can be solved by adjusting the input ionizing
continuum. Metallicity independent line ratios, C iii]/C iv
and N iii]1750/N iv]1846 suggests different ionization parameters
logU≃−1.5, and logU≃−2.5 for a range of gas column
densities. Since the ionization potential of N2+ (N+) is only
slightly larger than C2+ (C+), any smooth change in the continuum
slope in this regime is not likely to cause such big difference. However, we
note that the ionization potential (55.45eV) of N2+ is just above that of
He ii ionization potential, while that (47.89eV) of C2+ is below
the He ii ionization potential. If the continuum incident on the gas has
already been filtered by a relative highly ionized gas that is
optically thick to helium ionizing photons, such ionizing continuum
will have less N2+ ionizing photons relative to C2+ ionizing
photons, resulting more less N3+ ions. It would be interesting to
see if such a model can reproduce N iii]/N iv] and
C iii]/C iv ratio at the same time, but throughout
photo-ionization calculations are beyond the scope of this paper.
We only want to point out here that once such an ionizing continuum
is adopted, the ionization parameter will be close to, but
substantially larger than, the higher ionization paramter estimate
based on C iii]/C iv ratio.
Therefore, we interprete C4 as a dense outflow drivening by AGN.

4.1 Outflows

In Q 1321+058, we have detected emission lines from two outflow
components, C2 and C4. The broad component C4 with a radial velocity
1650 km s−1 has been interpreted as dense, metal enriched outflow
photoionized by AGN continuum, probably filtered by an ionized
absorber in last section. The narrow component C2 is also produced
by dense gas, but its nature is less clear. Therefore, we will focus
on the mass outflow rate and kinematic power of the high velocity outflow
C4 in this section first, and then discuss
briefly the C2 component.

We first estimate the mass in the outflow for C4. The H ii mass of the outflow
can be estimated from the Hα luminosity as follows

M=1.4mHLHαneαeffHαhνHα=450×(LHα1042ergs−1)(ne107cm−3)−1M⊙

(2)

assuming case-B recombination with a temperature of 104 K (Table 4.2
of Osterbrock 1989). The factor 1.4 accounts for Helium mass. With the
Hα luminosity of 1.5×1042 erg s−1, we obtain
∼680(ne/107)−1 M⊙ for the high velocity outflow.
This mass should be regarded as lower limits to the total mass in
consideration of the existence of possible HI regions. We
can estimate the mass of emission line gas using the best model
in section 3 as well. The model gives a similar value
∼400 M⊙.

For the single zone photoionization models discussed in §3.1, we estimate
the distance of the outflow to the AGN,

R=√Qion4πnHUc=150√107nH0.01Upc.

(3)

Assuming a typical quasar SED (Richards et al. 2006), the total ionizing
flux of Q 1321+058 is estimated to be 8×1057 photons s−1
(see §4.2). Using the density and ionization parameter
estimated in §3.1, we find distances of ∼80pc for C4, which is one
order of magnitude smaller than the size of NLR in quasars
of similar bolometric luminosity (Benert et al. 2005). Note the distance
is not very sensitive to the ionization parameter used.

The mass loss rate is small compared to the mass accretion rate onto
the supermassive black hole. The emission line gas is visible for a
dynamic time scale of ΔR/vr, where ΔR is the visible
radial length of the outflow, and vr is the observed radial velocity.
By combining with the mass in Eq 2, we obtain
an estimate of the mass loss rate,

Missing or unrecognized delimiter for \right

(4)

Note with the column density estimated in the section §3, the corrections due to HI gas are within a
factor of two. The mass loss rate is order
of 0.03 M⊙ yr−1 if ΔR/R∼1. Therefore, the mass
loss rate for these outflow is not important with respect to the
mass accretion rate onto the supermassive black hole (see next section).

The C4 outflow does not provide sufficient kinetic power to
account for the feedback that is supposed to heat the cold gas in
the host galaxy up to the escape velocity. The kinetic power
with the line emitting gas is only on the order of ∼10−6(100pc/ΔR) M⊙c2 yr−1, or 10−5(100pc/ΔR) of the AGN bolometric
luminosity according to the mass loss rate in Eq 4 and the observed
outflow velocity, while a few percent is required for the flow to be
energetically important (e.g., Di Matteo et al. 2005). A similar
conclusion has been reached by modelling absorption lines for nearby
Seyfert galaxies (Crenshaw, Kraemer & George 2003).

However, instead to conclude that outflows are energetically un-important
for AGN-galaxy feedback, we argue below that there is an indication for much
larger kinetic power than the above estimate.
As we have seen in the last section that the column density of the emission gas
is fairly low, of order of 1021 cm−2. A gas cloud with such column
density, density and typical temperature of photoionized gas is not bounded
by self-gravity, thus it will either disperse away on the time scales of only
10 years (i.e., ΔR∼2×10−2 pc)
or be confined by external pressure. In the former case, continuous formation
of the line emitting gas is required, and the mass outflow rate would be 40
M⊙ yr−1, and total kinetic power would be three magnitude higher than
the above estimate, which is sufficient to explain AGN-galaxy feedback. In the
latter case, the confined medium can be hot phase gas. Since the filling
factor of the warm gas is extremely small
(∼10−10×100pcΔR),
the bulk of the gas will be in the hot phase, and both the mass loss rate
and the kinetic power may be much larger than above estimate.
Hot gas has been detected in nearby
ultra-luminous infrared AGN including NGC 6240 (Komossa et
al. 2003). Future high spatial resolution X-ray observations can be
used to examine if such a hot X-ray component presents in this
quasar as well.

The covering factor of the emission line regions can be estimated from
the ratio of Hα to the ionizing photon flux assuming Case-B
recombination if the emission line region is optically thick to the
ionizing photons. It turns out that the covering factor of C4
is as small as 6×10−5. Similar numbers
(order of ∼10−4) are also obtained by comparing the C iii]1909,
[O iii]λ5007 or Hα luminosity with the best photoionization
model in the last section in combination with above distances. For
comparison, the covering factor of the NLR in quasars is typically a few
percent according to simple photon-counting estimation (e.g., Netzer
& Laor 1993). Thus, the AGN outflows in Q 1321+058 has a covering
factor 100 times smaller than a typical NLR.

In passing, with a maximum velocity extending to ∼4000 km s−1
and emission lines of a wide ionization range, Q1321+058 would appear
as a LoBAL quasar if the line of sight to the nuclear continuum is
intercepted by the outflow. We speculate that Q 1321+058 is an
obscured, high luminosity version of MKN 231, an extensively studied
infrared luminous low ionization BAL QSO with an extended outflow
region on kpc scale (e.g. (Lípari et al. 2005 )). The requirement of a
soft ionizing continuum for C4 is also consistent with weak
X-ray emission from BAL QSOs.

We would like to point out that the properties of outflow in C2 component
is very uncertain. First, Fe ii emission requires an extensive
warm HI region, but its column density cannot be properly constrained
without extensive modelling of the FeII emission. Thus, Eq (4)
will give only a lower limit to the mass loss rate. Second, if a large
fraction of Hα photons are produced through collision excitation
as is required to explain its large intrinsic Balmer decrement, Eq
(4), on the other side, Eq (4) is no long valid.
Finally, it has yet to be demonstrated that collisional excitation
can fully explain the large Hα/Hβ. Otherwise, dust
extinction may be important.

4.2 Quasar Activity and Star Formation History

The broad band spectral energy distribution (SED) of Q 1321+058 is
plotted in Fig. 6. The near-infrared data are taken from
the two micron all sky survey
(2MASS (Skrutskie et al. 2006)). The mid- and far-infrared data are
drawn from observations by the infrared space observatory (ISO)
((Kuraszkiewicz et al. 2003 )) and the IRAS faint source catalog ((Moshir et al. 1990 )),
respectively. Q1321+058 was also observed with Spitzer IRAC and
IRS instruments (Farrah et al. 2007; Sirocky et al. 2008).
The SDSS and HST FOS spectra, as well as the SDSS photometric data
corrected for Galactic extinction are also plotted.

The mid-infrared bump is very prominent in the SED, that rises
steeply from near- to mid-infrared, and then flattens and declines
in far-infrared. There is no apparent dip or peak around 9.7 micron
in the source rest frame. Lack of strong emission or absorption
feature is confirmed by Spitzer IRS observation, which reveals only
weak silicate absorption feature(Sirocky et al., 2008). The flat middle
to far-infrared spectrum and a lack of prominent PAH features
suggest that the bulk of the mid-infrared emission is powered by
the hidden quasar ((Weedman et al. 2005 ); (Hao et al. 2007)). To compare the observed
SED with that of red quasars, we redden the average SED of infrared
luminous quasar from Richards et al. (2006) through a uniform dust
screen with E(B−V)=4.5 using the extinction curve of Weingartner
& Draine (2001; 30ppm C in PAH), which can reproduce
approximately the extinction curve to the Galactic center. The
result is also shown in Fig. 3. This over-simplified model
can reproduce the overall shape of the infrared SED except for
predicting a strong dip around 10 μm due to small grains
(see also (Laor & Draine 1993 )), which can be eliminated with a detail
treatment of radiative transfer (e.g. (Siebenmorgen et al. 2004 )).

The 12μm to total observed [O iii] flux ratio is 50 time
higher than the average Seyfert galaxies (Haas et al. 2007). Even
after correcting for reddening, this ratio is still more than a factor of
ten higher. However, we consider that this can be attributed to
weak [O iii] emission rather than the dominance of mid-infrared
flux from the star-formation region for the aforementioned reason.
Note weak [O iii] emission is a common characteristics of low ionization
BAL QSOs (Boroson & Meyers 1993).

The total luminosity in the infrared is 1.9×1046 ergs s−1,
which makes Q 1321+058 an ULIR quasar. If Q 1321+058 bears an intrinsic
SED of a typical quasar, the bolometric luminosity is estimated to be
3.9×1046 ergs s−1, using the quasar template fit to
infrared SED and integrating in the optical to UV portion
3. However, if the
covering factor of
dust in this object is larger than the average value of quasars, the bolometric
luminosity is then over-estimated from the above fit. A
conservative estimate can be given by assuming that all of the UV
to X-ray continuum emission is absorbed and re-emitted isotropically
in the infrared. This yields a bolometric luminoisty a factor of two smaller.
The mass accretion rate is thus ∼ 3-6 M⊙ yr−1, for the typical
radiative efficiency of 0.1. If the quasar is accreting at close to the
Eddington limit, the black hole mass is estimated to be around 108 M⊙.

The SFR for this object may be estimated in several ways. We have
found that UV spectrum can be fitted by a young stellar population.
With the age and mass of the young stellar population, we estimate
a SFR of about 450-230 M⊙ yr−1. However, as we noted this
number is only valid to the order of magnitude giving our
oversimplified assumption about the model. Alternatively, using the UV
continuum luminosity, the SFR calibrator of Madau et al. (1998) and
assuming an extinction E(B-V)=0.26, we obtain a SFR of only
45 M⊙ yr−1. The discrepancy arises because the SSP model contains
1 Myr population only, while older stellar populations still
contribute to the observed UV flux. The SFR can also be estimated from
the far-infrared luminosity and radio power. Assuming most of the far-infrared
luminosity is powered by the starburst(e.g., (Schweitzer et al. 2006 )), we find an
upper limit to SFR of about 270 M⊙ yr−1, similar to that seen in
NGC 6240 (e.g., Pasquali, Gallagher & de Grijs 2004). Q 1321+058 was
detected in both NVSS and FIRST with a flux density of
S1.4GHz,NVSS=4.9±0.5 mJy and S1.4GHz,FIRST=4.5±0.3 mJy.
Assuming that the radio emission (6×1030ergs−1Hz−1) all
comes from the H ii region, we obtain an upper limit of SFR≲330M⊙yr−1. These values are consistent with that derived from the UV
spectral fit. The SFR can be also estimated with the PAH luminosity
(Wu et al., 2005, e.g.). (Cao et al. 2008) detected a 6.2μm PAH luminosity
of 1.2×1042 erg s−1. Following Hernan-Caballero et al.
(2009), we obtain a SFR of only 4.6 M⊙ yr−1. Finally, the SFR can be
estimated from Hα luminosity. An upper limit can be derived if
C1 is assumed to be ionized by young stars. The extinction corrected
Hα luminosity for C1 are 1.8×1042 erg s−1.
According to the calibration of Madau et al. (1998), the SFR is
14 M⊙ yr−1. SFR will be one order of magnitude larger if C2 is
from star-formation region.
The very different SFR derived from different methods
demonstrate the difficulty of estimating SFR in the presence of AGN activity.

With all above estimates, SFR does not likely exceed hundreds M⊙ yr−1
, while we have found a mass accretion rate of 3-6 M⊙ yr−1 to the black hole according to its bolometric
luminosity. The ˙M/SFR is more than an order of magnitude larger
than that required for the strict co-evolution of the black hole and
bulge. However, the black hole to stellar mass ratio can still be
consistent with that found in local spheroidal galaxies
(Gebhardt et al. 2000). The mass of the intermediate-aged stellar
population in Q 1321+058 is ∼6×1010 M⊙ from our
analysis of the SDSS and HST
spectra. If the bolometric luminosity does not exceed the Eddington
luminosity, a lower limit to the black hole can be set to 108 M⊙.
It was argued that ultraluminous infrared quasars are accreted at
the Eddington limit rather than fuel limited (Hao et al. 2005).
Then a lower-limit on a mass ratio of the black hole to the intermediate
and young stellar populations to be ∼2×10−3, in coincidence
with that for local spheroidal galaxies (Gebhardt et al. 2000).

We have performed a detailed analysis of the optical–UV
spectrum and the broad band spectral energy distribution
of Q 1321+058. Our analysis confirms that it is an obscured quasar
with its bulk energy output in mid-infrared.
Its optical–UV spectra show complex emission lines.
We identified four components: a narrow component at the systematic
velocity (C1), a narrow component at velocity -380 km s−1 (C2), and
two broad components at -80km s−1 (C3) and 1650km s−1 (C4), respectively.
C1 shows a LINER/composite spectrum, which is common among ULIRGs. Both
C2 and C4 can be interpreted as dense outflows with the back-sides are
obscured. We speculate that C3 comes from an intermediate line region
between BLR and NLR.

A comparison of the measured emission line ratios with photo-ionization
models suggests that C4 outflow has a gas density nH∼107 cm−3,
column density NH∼1021 cm−2, an α-enriched
super-solar metallicity of Z∼10Z⊙ for starburst galaxies. It
is located at a distance about a hundred parsecs from the central
continuum source. The velocity range, ionization level, and column
density derived from the emission lines suggest that Q 1321+058 might
be a low ionization broad absorption line (LoBAL) quasar viewed in an
“unfavored” direction, with the outflows being part of the
otherwise LoBAL region. The apparent total mass loss rate in
the C4 outflow is small, thus the outflow associated with line emitting
gas does not have sufficient energy to remove the ISM of host galaxies,
quenching both star-formation and accretion process.
But we argued that the optical emission line region may trace
only a small amount of quasar outflows, thus the actual mass loss rate
and kinetic power may be much larger. The covering factor of the outflow
is very small.

The optical and UV continuum can be well modeled with a young (∼1Myr) plus an intermediate age (0.5−1 Gyr) stellar population.
The latter population has a mass around a few 1010
M⊙, suggesting for a recent building of a massive galaxy
following the merger of two gas rich galaxies. Fast
stellar mass building is also consistent with the metallicity that
required to explain the C3 and C4 line ratios. We obtain very different
SFR or its upper limits, independently from UV continuum, radio,
far-infrared and emission line luminosity in the range of a few to
serval hundred M⊙ yr−1. We estimate a bolometric luminosity
of (2-4) 1046 erg s−1 for the quasar, or a mass accretion rate of
3-6 M⊙ yr−1 for a typical
efficiency of 0.1. The SFR to mass accretion ratio is more than one order
of magnitude lower than that is required for the co-evolution of black
hole and spheroid. If the black hole is accreted at the Eddington
rate, the black mass to the stellar mass ratio will be coincident with
that defined by local spheroid galaxies.

We thank the anonymous referee for constructive comments which lead
to significant improvement in the presentation. This work is supported
by Chinese Natural Science Foundation through CNSF-10233030 and
CNSF-10573015, and by the Knowledge Innovation Program of the Chinese
Academy of Sciences, Grant No. KJCX2-YW-T05. H. Y. Zhou acknowledges
the Chinese NSF support through NSFC-10473013, and the support from
NSF AST-0451407 and AST-0451408, NASA NNG05G321G and NNG05GR41G, and
the University of Florida.

Appendix A Notes on the fit to [O iii], N iii] and Si iii] emission lines

Initially, we fit the four components to the
[O iii]λλ4959,5007 blend, and for each component the
doublet ratio is fixed to 1/3. The fit is not satisfactory. In
particular, there is an excess in the blue component in
[O iii]λ4959. This indicates that the doublet ratio of one or
more components is not the theoretical value, or there is
contribution from other emission line(s). Potential contaminating
lines in the regime can be [Fe vii]λ4989, He iλ5016,
[Fe iiI]λλ4988,4931 or the Fe ii blend. Since the Fe ii bumps
around 4700 Å and 5100 Å is not visible, we will not consider
this as a likely possibility4. Our grid photo-ionization
calculation suggests that the He i emission should be no more than
2% of Hβ in strength
(for a wide range of parameters as discussed in Appendix 2),
and is too weak to account for the excess. [Fe vii] λ4989
can be fairly strong for a high ionization parameter and a low column
density in photoionization models, but [Fe vii]λ6087, which
is not detected in our spectrum, should be a factor of at least 2.7
stronger for all models.
Therefore, it is likely that the excess is
due to the [O iii]λ4959 emission. To incorporate this, we leave
the [O iii]λ4959/λ5007 ratio of the blue component as
a free parameter. This gives fairly good fit with the ratio of 0.46.
Therefore, in the subsequent fit, the [O iii] doublet ratio is
allowed to vary freely for C4 and is fixed at 1/3 for other
components.

The N iii]λ1750 emission is a blend of five lines, and
the multiplets ratios depend on both gas temperature and density
((Dwivedi et al. 1995 )). Using the grid photo-ionization models described in
Appendix 2, we estimate that this feature is dominated by
N iii]λ1751 and N iii]λ1752, and other
lines of the multiplets account for less than 10% only. The line
ratios are only weakly depends on the ionization parameters, we
fixed the multiple ratios at the model value in the vicinity of
the best model density.

Both Si iii]λ1889 and [Si iii]λ1883 may
contribute to the λ1888 blend. At a low density, [Si iii]
is the main contributor
to this feature, while at a density above 106 cm−3 this
feature is dominated by Si iii]. The situation is similar for the
λ1909 blend of [C iii]λ1907 and C iii]λ1909.
As aforementioned, the Si iii/C iii ratio requires this
feature being dominated by semi-forbidden transitions. Therefore,
we take wavelengths 1889Å and 1909Å for Si iii and C iii, respectively,
in the following analysis.

Appendix B Photoionization Models

We have computed large grid constant-density models using the
photoionization code CLOUDY (version 06.02, last described by
Ferland et al. 1998). The range of gas density considered is 102−10cm−3, and the range of ionization parameter considered is
−3.5<logU<0.0. We consider five column densities of logNH[cm−2]=21, 21.5, 22.0, 22.5, and 23.0, and four metal
abundance values: 1, 3, 5 and 10 Z⊙. Two metal enrichment
schemes are considered: all metal abundances are scaled-up solar
values, and metal enrichment in a starburst (Haman & Ferland 1997).
We choose a typical AGN ionizing continuum (Korista et al. 1997).
We examine the line ratio contours on the dual-parameter planes
for the plausible regimes (an example of such contours is shown
in Fig 7). We find that: (1) with a solar and
scaled solar metallicity, large Si iii]/C iii]∼1 can be
reproduced only at a high density nH>109.5 cm−3 for
the parameter ranges explored, while it can be produced in a wide
parameter range for the starburst metallicity; (2) a large N iii]/C iii]
ratio for C4 can only be reproduced at a large starburst metallicity
Z>12Z⊙. At small column densities (logNH<22), the line
ratio is high at a relative high ionization parameters; while at
large column density, the line ratio is essentially independent of
ionization parameters; (3) for C4, in the column density and density
ranges considered here, N iii]/N iv] suggests a low ionization parameter
than C iii]/C iv.

Figure 2: Upper panel: The two-stellar population model for the optical
spectrum. The UV model with the SMC (Calzetti et al’s) extinction curve is
shown in green (brown) color, and the intermediate
age stellar population in yellow, and the final model in red.
Lower panel: The UV spectrum
and the best-fit reddened single stellar population model (the SMC
extinction curve; Green curves).
The fit to the UV continuum and emission lines is shown in red.

Figure 3: Continuum subtracted emission line profiles and the best-fit
four component models, green for C1, red for C2, purple for C3, and yellow
for C4 (see text for details).
Figure 4: BPT diagrams for C1 (black diamonds) and C2 (blue diamonds).
Black curve represents the empirical boundary separating star-forming
galaxies from AGNs (Kauffmann et al. 2003), while red curves for the
theoretical curve of extreme starburst (Kewley et al. 2001). The blue
straight lines divide LINERs from Seyfert galaxies.Figure 5: Wavelength-dependent Hα/Hβ ratio over the
emission line profile. The best-fit [N ii] model has been subtracted
from the Hα+[N ii] blend. It shows a clear peak at position of
C2. The horizon line marks Hα/Hβ=3.1.
Figure 6: Spectral energy distribution from far-infrared to ultraviolet
for Q1321+058. The infrared SED of NGC 6240 is shown in solid line, the
template of infrared-luminous quasars in Richards et al. (2006)
in red, while the quasar template reddened by E(B-V)=4.5 in blue.
Figure 7: Contours of emission line ratios in the ionization parameter
– density space calculated using the Cloudy C06.02 (Ferland 1998) for
metal abundance Z=10Z⊙and a column density NH=1021 cm−2.
The regimes for C4 is between two colored lines (green and blue) after
taking account for the one sigma error bar for the observed line ratios.

Footnotes

slugcomment: submitted to ApJ,March 7, 2018

http://www.sdss.org/.

Infrared emission is the dust reprocessed light from the optical
and ultraviolet continuum, thus should not be counted in the bolometric
luminosity calculation for an unobscured quasar.